1 - Astronomical Observatory of Odessa National University and Isaac Newton Institute of Chile, Shevchenko Park, 65014 Odessa, Ukraine
2 - Observatoire de Bordeaux, CNRS UMR 5804, BP 89, 33270 Floirac, France
3 - Steward Observatory, The University of Arizona, Tucson, AZ 85721, USA

Abstract
Line depth ratios measured on high resolution (R=42 000), high S/N echelle spectra are used for the determination of precise effective temperatures of 181 F, G, K Main Sequence stars with about solar metallicity ( $-0.5 < [{\rm Fe/H]} < +0.5$ ). A set of 105 relations is obtained which rely $T_{\rm eff}$ on ratios of the strengths of lines with high and low excitation potentials, calibrated against previously published precise (one per cent) temperature estimates. The application range of the calibrations is 4000-6150 K (F8V-K7V). The internal error of a single calibration is less than 100 K, while the combination of all calibrations for a spectrum of S/N=100reduces uncertainty to only 5-10 K, and for S/N=200 or higher - to better than 5 K. The zero point of the temperature scale is directly defined from reflection spectra of the Sun with an uncertainty about 1 K. The application of this method to investigation of the planet host stars properties is discussed.

The direct method to determine the effective temperature of a star relies on the measurement of its angular diameter and bolometric flux. In practice certain limitations restrict the use of this fundamental method to very few dwarfs. Other methods of temperature determination have errors of the order 50-150 K, which translates into the [Fe/H] error of $\sim$ 0.1 dex or larger. The only technique capable so far of increasing this precision by one order of magnitude, is the one employing ratios of lines with different excitation potentials $\chi$ . As is well known, the lines of low and high $\chi$ respond differently to the change in $T_{\rm eff}$ . Therefore, the ratio of their depths $r=R_{\lambda1}/R_{\lambda2}$ (or equivalent widths, EW) should be very sensitive temperature indicator. The big advantage of using line-depth ratios is the independence on the interstellar reddening, spectral resolution, rotational and microturbulence broadening.

The reader is referred to Gray (1989, 1994) and Gray & Johanson (1991) to learn more about the history and justification of the line ratio method. Applying this method to the Main-Sequence (MS) stars, they achieved precision as high as 10 K. The most recent works on the subject are by Caccin et al. (2002) who discuss the possible weak points of this technique for the case of dwarfs (see below), and the fundamental contribution by Strassmeier & Schordan (2000) who report 12 temperature calibrations for giants with an error of only 33 K.

So far however the line-ratio method has not been fully utilized for purposes other than just temperature estimation by itself. One of few applications is the chemical abundance analysis of supergiants, where it has proved the anticipated high efficiency and accuracy. Thus, Kovtyukh & Gorlova (2000, hereafter Paper I) using high-dispersion spectra, established 37 calibrations for the temperature determination in supergiants (a further study increased this number to 55 calibrations). Based on this technique, in a series of 3 papers, Andrievsky et al. (2002, and references therein) derived temperatures for 116 Cepheids (from about 260 spectra) at a wide range of galactocentric distances ( $R_{\rm g}=5{-}15$ kpc) with a typical error 5-20 K. The high precision of this new method of temperature determination allowed them to uncover the fine structure in the Galactic abundance gradients for many elements. Even for the most distant and faint objects ( $V \simeq 13{-}14$ mag) the mean error in $T_{\rm eff}$ was no larger than 50-100 K, with maximum of 200 K for spectra with lowest S/N (=40-50).

Another example concerns T Tau stars. For young stars, uncertainties in reddening due to variable circumstellar extinction invalidate the photometric color method of effective temperature determination. Using 5 ratios of FeI and VI lines calibrated against 13 spectral standards, Padgett (1996) determined the effective temperature of 30 T Tau stars with a 1 $\sigma$ uncertainty lower than 200 K.

The intent of this paper is to improve this technique, based on our experience of applying it to supergiants (Paper I and following publications), and expand it to the MS stars. The wide spectral range of ELODIE echelle spectra allowed to select many unblended lines of low and high excitation potentials thus improving the internal consistency of the method, whereas the large intersection between the ELODIE database and published catalogues of effective temperatures allowed to take care of systematic effects. We obtained a median precision of 6 K on $T_{\rm eff}$ derived for an individual star. The zero-point of the scale was directly adjusted to the Sun, based on 11 solar reflection spectra taken with ELODIE, leading to the uncertainty in the zero-point of about 1 K.

Temperature determined by the line ratio method may now be considered as one of the few fundamental stellar parameters that have been measured with an internal precision of better than 0.2%.

2 Observations and temperature calibrations

The investigated spectra are part of the library collected with the ELODIE spectrometer on the 1.93-m telescope at the Haute-Provence Observatory (Soubiran et al. 1998; Prugniel & Soubiran 2001). The spectral range is 4400-6800 Å and the resolution is R=42 000. The initial data reduction is described in Katz et al. (1998). All the spectra are parametrized in terms of $T_{\rm eff}$ , $\log g$ , [Fe/H], either collected from the literature or estimated with the automated procedure TGMET (Katz et al. 1998). This allowed us to select a sample of spectra of FGK dwarfs in the metallicity range $-0.5 < [{\rm Fe/H]} < +0.5$ . Accurate Hipparcos parallaxes are available for all of the stars of interest enabling to determine their absolute magnitudes M_V that range between 2.945 (HD 81809, G2V) and 8.228 (HD 201092, K7V). All the selected spectra have a signal to noise ratio greater than 100 (see Fig. 1). Further processing of spectra (continuum placement, measuring equivalent widths, etc.) was carried out by us using the DECH20 software (Galazutdinov 1992). Equivalent widths EWs and depths $R_{\lambda}$ of lines were measured manually by means of a Gaussian fitting. The Gaussian height was then a measure of the line depth. This method produces line depths values that agree nicely with the parabola technique adopted in Gray (1994). We refer the reader to Gray (1994, and references therein) and Strassmeier & Schordan (2000) for a detailed analysis of error statistics.

Following Caccin's et al. (2002) results, where a careful analysis of the anticipated problems for the Solar-type stars has been carried out, we did not use ion lines and high-ionization elements (like C, N, O) due to their strong sensitivity to gravity.

Gray (1994) showed that the ratio of lines VI 6251.82 and FeI 6252.55 depends strongly on metallicity. The reason is that the strong lines like FeI 6252.55 ( $R_{\lambda}=0.52$ for the Sun) are already in the dumping regime, where the linearity of EW on abundance breaks down. In addition, as was shown in the careful numerical simulations by Stift & Strassmeier (1995), this ratio (of 6251.82 and 6252.55 lines) is also sensitive to rotational broadening. Significant effects were found for $v\sin i$ as small as 0-6 km s^-1 (for solar-like stars). We therefore avoided using strong lines in our calibrations. Indeed, Gray (1994) concluded that, as expected, the weak-line ratios are free from the effects of metallicity. As to the effect of rotation, we should note that all objects in our sample are old Main Sequence stars with slow to negligible rotation ( $v\sin i<15$ km s^-1), which is comparable to the instrumental broadening.

Thus, we initially selected about 600 pairs of 256 unblended SiI, TiI, VI, CrI, FeI, NiI lines with high and low excitation potentials within the wavelength interval 5300-6800 Å.

(1) the excitation potentials of the lines in a pair must differ as much as possible;

**Table 1:** Program stars. Asterisks indicate stars with planets.
HD/BD	HR	Name	$T_{\rm eff}$	N	$\sigma$ , K	$T_{\rm eff}$	$T_{\rm eff}$	$T_{\rm eff}$	$T_{\rm eff}$	Mv	B-V	rem
			this paper			EDV93	AAMR96	BLG98	DB98
1562	-		5828	97	5.8					5.006	0.585
1835	88	9 Cet	5790	68	5.5			5713	5774	4.842	0.621
3765	-		5079	87	4.7					6.161	0.954
4307	203	18 Cet	5889	91	5.0	5809	5753		5771	3.637	0.568
4614	219	24 Eta Cas	5965	69	6.4	5946	5817			4.588	0.530
5294	-		5779	86	6.6					5.065	0.610
6715	-		5652	97	6.7					5.079	0.658
8574	-		6028	61	6.7					3.981	0.535	*

(2) the lines must be close in wavelength; it turned out though that calibrations based on widely spaced lines (including from different orders) show same small dispersion as the closely spaced lines. Therefore, we retained all pairs with a difference in wavelength up to 70 Å ( $\lambda_{2} - \lambda_{1}<70$ Å);

(3) the lines must be weak enough to eliminate a possible dependence on microturbulence, rotation and metallicity;

(4) the lines must be situated in the spectral regions free from telluric absorption.

The next step was to choose the initial temperatures for interpolation. This is a very important procedure since it affects the accuracy of the final temperature scale, namely, the run of the systematic error with $T_{\rm eff}$ (Fig. 3). There is an extended literature on MS stars temperatures. For 45 stars from our sample (see Table 1) we based the initial temperature estimates on the following 3 papers: Alonso et al. (1996, hereafter AAMR96), Blackwell & Lynas-Gray (1998, hereafter BLG98) and DiBenedetto (1998, hereafter DB98). In these works the temperatures have been determined for a large fraction of stars from our sample with a precision of about 1%. AAMR96 used the Infrared Flux Method (IRFM) to determine $T_{\rm eff}$ for 475 dwarfs and subdwarfs with a mean accuracy of about 1.5% (i.e., 75-90 K). BLG98 also have determined temperatures for 420 stars with spectral types between A0 and K3 by using IRFM and achieved an accuracy of 0.9%. DB98 derived $T_{\rm eff}$ for 537 dwarfs and giants by the empirical method of surface brightness and Johnson broadband (V-K) color, the accuracy claimed is $\pm$ 1%. Whenever 2 or 3 estimates were available for a given star, we averaged them with equal weights. These temperatures served as the initial approximations for our calibrations.

First, for the mentioned above 45 stars with previously accurately determined $T_{\rm eff}$ we plotted each line ratio against $T_{\rm eff}$ , and retained only those pairs of lines that showed unambiguous and tight correlation. We experimented with a total of nearly 600 line ratios but adopted only the 105 best - the ones showing the least scatter. These 105 calibrations consist of 92 lines, 45 with low ( $\chi< 2.77$ eV) and 47 with high ( $\chi>4.08$ eV ) excitation potentials. Judging by the small scatter in our final calibrations (Fig. 2) and $T_{\rm eff}$ , the selected combinations are only weakly sensitive to effects like rotation, metallicity and microturbulence. This confidence is reinforced by the fact that the employed lines belong to a wide range of chemical elements, intensity and atomic parameters, therefore one can expect the mutual cancellation of opposite effects.

Each relationship was then fitted with a simple analytical function. Often calibrations show breaks which not can be adequately described even by a 5th-order polynomial function (see Fig. 2). Therefore, we employed other functions as well, like the Hoerl function ( $T_{\rm eff}=ab^{r}*r^{c}$ , where $r=R_{\lambda1}/R_{\lambda2}$ , a,b,c - constants), modified Hoerl ( $T_{\rm eff} =ab^{1/r} r^{c}$ ), power low ( $T_{\rm eff}=ar^{b}$ ), exponential ( $T_{\rm eff}=ab^{r}$ ) and logarithmic ( $T_{\rm eff}=a{+}b~ln(r)$ ) functions. For each calibration we selected the function that produced the least square deviation. As a result, we managed to accurately approximate the observed relationships with a small set of analytic expressions. This first step allowed us to select 105 combinations, with an rms of the fit lower than 130 K, the median rms being 93 K. Using these initial rough calibrations, for each of the 181 target stars we derived a set of temperatures (70-100 values, depending on the number of line ratios used), averaged them with equal weights, and plotted these mean $T_{\rm eff}$ (with errors of only 10-20 K) versus line ratios again, thus determining the preliminary calibrations (for which the zero-point had yet to be adjusted).

We would like to point out that the precision of our calibrations varies with temperature. In particular, at high $T_{\rm eff}$ the lines with low $\chi$ become very weak causing the line depth measurement to be highly uncertain. Therefore, for each calibration we determined the optimum temperature range where the maximum accuracy is attained (no worse than 100 K), so that for a given star only a subset of calibrations can be applied.

What are the main sources of random errors in the line ratio method? The measurement errors in line depths are mainly caused by the continuum placement uncertainty and by the Gaussian approximation of the line profile. In addition, the individual properties of the stars, such as metallicity, spots, rotation, convection, non-LTE effects and binarity may also be responsible for the scatter observable in Fig. 2. The detailed analysis of these and other effects can be found in Paper I, Strassmeier & Schordan (2000) and in works by D. F. Gray. We estimate that the typical error in the line depth measurement $r=R_{\lambda1}/R_{\lambda2}$ is 0.02-0.05, implying an error in temperature of about 20-50 K.

The mean random error of a single calibration is 60-70 K (40-45 K in the most and 90-95 K in the least accurate cases). The use of $\sim$ 70-100 calibrations reduces the uncertainty to 5-7 K (for spectra with S/N=100-150). Better quality spectra (R>100 000, S/N>400) should in principle allow an uncertainty of just 1-2 K. Clearly, time variation of the temperature for a given star should be readily detected by this method, since the main parameters that cause scatter due to star-to-star dissimilarities (gravity, rotation, [Fe/H], convection, non-LTE effects etc.) are fixed for a given star. The temperature variation of several degrees in mildly active stars may be produced by the surface features and rotational modulation, as for example has been documented for the G8 dwarf $\xi$ Bootis A (Toner & Gray 1988) and $\sigma$ Dra (K0V, Gray et al. 1992).

The next stage is to define the zero point of our temperature scale. Fortunately, for dwarfs (unlike for supergiants) a well-calibrated standard exists, the Sun. Using our preliminary calibrations and 11 independent solar spectra from the ELODIE library (reflection spectra of the Moon and asteroids), we obtained a mean value of $5733 \pm 0.9$ K for the Sun's temperature. Considering the Sun as a normal star, we adjusted our calibrations by adding 44 K to account for the offset between the canonical Solar temperature of 5777 K and our estimate. The possible reasons for this small discrepancy are discussed below.

3 Results and discussion

Table 1 contains our final $T_{\rm eff}$ determinations for 181 MS stars. Note that we added the 44 K correction to the initial calibrations in order to reproduce the standard 5777 K temperature of the Sun. For each star we report the mean $T_{\rm eff}$ , number of the calibrations used (N), and the standard error of the mean ( $\sigma$ ). For comparison, we also provide $T_{\rm eff}$ as determined in Edvardsson et al. (1993, hereafter EDV93), AAMR96, BLG98 and DB98. Absolute magnitudes M_V have been computed from Hipparcos parallaxes and V magnitudes from the Tycho2 catalogue (Høg et al. 2000) transformed into Johnson system. (B-V) are also from Tycho2. Planet-harboring stars are marked with an asterisk.

As one can see from Table 1, for the majority of stars we get an error which is smaller than 10 K. The consistency of the results derived from the ratios of lines representing different elements is very reassuring. It shows that our 105 calibrations are essentially independent of micro-turbulence, LTE departures, abundances, rotation and other individual properties of stars. We admit though that a small systematic error may exist for $T_{\rm eff}$ below 5000 K where we had only few standard stars.

As was already mentioned, for the first approximation we took accurate temperatures from AAMR96, BLG98 and DB98. The comparison of our final $T_{\rm eff}$ with those derived by AAMR96, BLG98 and DB98 is shown in Fig. 3. As a test of the internal precision of our $T_{\rm eff}$ we investigate the $T_{\rm eff}$ - color relation with the Strömgren index b-y, using our determinations of $T_{\rm eff}$ , and those obtained by other authors. The results are shown in Table 2 where the rms of the linear fit is given for each author's determination, along with our estimate of $T_{\rm eff}$ and using common stars. In each case the scatter of the color relation is significantly improved when adopting our temperatures, though some residual dispersion is still present that can be attributed to the photometric errors, reddening and the intrinsic properties of stars (metallicity, gravity, binarity...) to which the color indices are known to be sensitive. The improvement is particulary spectacular in comparison with EDV93. This proves the high quality of our temperatures and the mediocrity of b-y as a temperature indicator.

**Table 2:** RMS of a linear regression between $T_{\rm eff}$ and Strömgren b-y using effective temperatures obtained by other authors and in this study with N common stars.
author	N	$\sigma_{\rm others}$ K	$\sigma_{\rm our}$ K
AAMR96	30	102	78
BLG98	25	71	65
DB98	29	113	87
EDV93	30	63	29

Another point concerns the difference between the zero-point of our temperature scale and that of other authors. Comparing 30 common objects, we find that the AAMR96 scale underestimates temperatures by 45 K near the solar value compared to ours, but apart from that, the deviations are random and no trend with $T_{\rm eff}$ is present. The 45 K offset may arise from the various complications associated with observing the Sun as a star, and/or problems in the models used by AAMR96, like underestimation of convection in the grid of the model atmosphere flux developed by Kurucz. After correcting the AAMR96 zero-point for the 45 K offset, the mean random error of their scale becomes 65 K (where we neglect the error of our own scale which is an order of magnitude less).

The temperatures of BLG98 are also in a good agreement with our estimates - except for a 48 K offset, no correlation of the difference with temperature is observed. The mean dispersion is 63 K (for 26 common stars), which is within the errors of the BLG98 scale.

Comparing with DB98: for the 29 stars in common, their temperatures are on average 41 K below ours, and the mean error is $\pm$ 53 K.

Thus, the temperatures derived in AAMR96, BLG98 and DB98 have good precision, though the absolute values are somewhat low relative to the Sun. The reason may be due to the difficulty of photometric measurements of the Sun, as well as indicating some problems in the model atmosphere calculations employed. For example, the Sun's temperatures derived in AAMR96 and DB98 are identical - 5763 K, which is below the nominal value of 5777 K. Besides, the mean temperatures of solar analogue stars (spectral types G2-G3, $\rm [Fe/H]\approx 0.0$ , and Sun being of type G2.5) derived in these papers are significantly below the solar value: $5720\pm54$ K (AAMR96, 3 stars), $5692 \pm31$ K (BLG98, 11 stars) and $5702\pm46$ K (DB98, 7 stars). Our determination for the G2-G3 spectral types is $5787\pm14$ K, based on 12 stars. This demonstrates that a small error (0.8%) affects the zero point of the IRFM method, because when applied to the Sun and the solar type stars, it returns inconsistent results.

We also compared our estimates of $T_{\rm eff}$ with photometric temperatures. EDV93 derived temperatures of 189 nearby field F, G disk dwarfs using the theoretical calibration of temperature versus Strömgren (b-y) photometry (see Table 1). The mean difference between the $T_{\rm eff}$ of Edvardsson et al. (1993) and ours is only -14 K ( $\sigma=\pm$ 67 K, based on 30 common stars).

To compare our temperatures to Gray (1994), we used his calibration of $(B-V)_{\rm corr}$ corrected for metallicity. Our scale is +11 K lower ( $\sigma=\pm$ 61 K, 24 stars).

Summarizing, we demonstrated that our temperature scale is in excellent agreement with the widely used photometric scales, while both the IRFM method and the method of surface brightness predict too low values for the temperature of the Sun and the solar type stars.

Figure 4 shows the sensitivity of our technique to temperature. Two outliers with errors greater than 20 K are the cold dwarfs HD 28343 and HD 201092, known as flaring stars. For other stars the internal errors range between 3 and 13 K, with a median of 6 K.

4 Conclusion

The high-precision temperatures were derived for a set of 181 dwarfs, which may serve as temperature standards in the 4000-6150 K range. These temperatures are precise to within 3-13 K (median 6 K) for the major fraction of the sample, except for the two outliers. We demonstrated that the line ratio technique is capable of detecting variations in $T_{\rm eff}$ of a given star as small as 1-5 K. This precision may be enough to detect star spots and Solar-type activity cycles. Of particular interest is the application of this method to testing ambiguous cases of low-mass planet detection, since planets do not cause temperature variations, unlike spots.

The next step will be the adaptation of this method to a wider range of spectral types and for an automatic pipeline analysis of large spectral databases.

References

5 Online Material

Table 1F: Program stars. Asterisks indicate stars with planets.
HD HR Name $T_{\rm eff}$ N $\sigma$ , K $T_{\rm eff}$ $T_{\rm eff}$ $T_{\rm eff}$ $T_{\rm eff}$ Mv B-V rem

this paper EDV93 AAMR96 BLG98 DB98

1562 - 5828 97 5.8 5.006 0.585

1835 88 9 Cet 5790 68 5.5 5713 5774 4.842 0.621

3765 - 5079 87 4.7 6.161 0.954

4307 203 18 Cet 5889 91 5.0 5809 5753 5771 3.637 0.568

4614 219 24 Eta Cas 5965 69 6.4 5946 5817 4.588 0.530

5294 - 5779 86 6.6 5.065 0.610

6715 - 5652 97 6.7 5.079 0.658

8574 - 6028 61 6.7 3.981 0.535 *

8648 - 5790 59 7.2 4.421 0.643

9826 458 50 Ups And 6074 44 13.1 6212 6155 6136 3.452 0.496 *

10145 - 5673 96 4.2 4.871 0.667

10307 483 5881 94 4.0 5898 5874 4.457 0.575

10476 493 107 Psc 5242 69 3.2 5172 5223 5157 5.884 0.819

10780 511 5407 95 4.0 5.634 0.767

11007 523 5980 84 7.4 3.612 0.524

13403 - 5724 91 7.0 5585 5577 5588 3.949 0.616

13507 - 5714 91 5.4 5.123 0.637 *

13825 - 5705 96 5.5 4.700 0.674

14374 - 5449 77 4.6 5.492 0.757

15335 720 13 Tri 5937 84 6.6 5857 5921 3.468 0.539

17674 - 5909 58 8.7 5875 5880 4.194 0.563

17925 857 5225 87 5.0 5.972 0.864

18803 - 5659 95 3.5 4.998 0.669

19019 - 6063 56 7.2 4.445 0.508

19308 - 5844 95 5.4 4.220 0.626

19373 937 Iot Per 5963 75 5.1 5996 5981 5951 3.935 0.554

19994 962 94 Cet 6055 56 10.0 6104 3.313 0.523 *

22049 1084 18 Eps Eri 5084 84 5.9 5076 6.183 0.877 *

22484 1101 10 Tau 6037 60 3.6 5981 5998 5944 5940 3.610 0.527

23050 - 5929 80 9.0 4.330 0.544

24053 - 5723 93 3.7 5.183 0.674

24206 - 5633 94 4.8 5.418 0.681

26923 1322 V774 Tau 5933 77 5.9 4.685 0.537

28005 - 5980 87 6.1 4.359 0.652

28099 - 5778 85 5.2 4.747 0.660

28343 - 4284 20 20.3 8.055 1.363

28447 - 5639 93 6.3 3.529 0.678

29150 - 5733 89 5.4 4.934 0.668

29310 - 5852 89 7.7 5781 5775 4.407 0.564

29645 1489 6009 57 5.8 6028 3.504 0.548

29697 - 4454 40 11.4 7.483 1.108

30495 1532 58 Eri 5820 91 5.7 4.874 0.588

30562 1536 5859 87 6.8 5886 5822 5843 5871 3.656 0.593

32147 1614 4945 65 8.7 6.506 1.077

34411 1729 15 Lam Aur 5890 88 4.3 5889 5847 5848 5859 4.190 0.575

38858 2007 5776 81 6.7 5669 5697 5.014 0.584

39587 2047 54 Chi1 Ori 5955 71 6.1 5953 4.716 0.545

40616 - 5881 89 10.0 3.833 0.585

41330 2141 5904 77 5.5 5917 4.021 0.547

41593 - 5312 92 3.3 5.814 0.802

42618 - 5775 96 6.6 5.053 0.603

42807 2208 5737 81 5.2 5.144 0.631

43587 2251 5927 81 4.4 4.280 0.558

43947 - 6001 82 7.1 5945 4.426 0.507

45067 2313 6058 61 4.6 3.278 0.507

47309 - 5791 95 3.9 4.469 0.623

50281 - 4712 56 8.5 6.893 1.074

50554 - 5977 77 5.8 4.397 0.529 *

51419 - 5746 94 8.3 5.013 0.600

55575 2721 5949 65 6.6 5963 5839 4.418 0.531

58595 - 5707 87 8.3 5.105 0.665

60408 - 5463 97 4.7 3.100 0.760

61606 - 4956 83 4.6 6.434 0.955

62613 2997 5541 90 6.4 5.398 0.695

64815 - 5864 88 8.3 3.375 0.605

65874 - 5936 85 4.7 3.100 0.574

68017 - 5651 100 9.0 5512 5.108 0.630

68638 - 5430 90 6.3 5.021 0.746

70923 - 5986 82 4.5 3.879 0.556

71148 3309 5850 88 5.1 4.637 0.587

72760 - 5349 91 3.8 5.628 0.796

72905 3391 3 Pi1 UMa 5884 79 6.8 4.869 0.573

73344 - 6060 37 6.8 4.169 0.515

75318 - 5450 78 5.8 5.345 0.717

75732 3522 55 Rho1 Cnc 5373 97 9.7 5.456 0.851 *

76151 3538 5776 88 3.0 5763 4.838 0.632

76780 - 5761 87 5.0 5.011 0.648

81809 3750 5782 85 6.9 5611 5619 2.945 0.606

82106 - 4827 76 6.0 6.709 1.000

86728 3951 20 LMi 5735 91 5.6 5746 4.518 0.633

88072 - 5778 82 5.0 4.717 0.593

89251 - 5886 89 6.3 3.292 0.569

89269 - 5674 95 5.7 5.089 0.645

89389 4051 6031 48 8.9 4.034 0.532

91347 - 5923 75 7.4 4.725 0.513

95128 4277 47 UMa 5887 89 3.8 5882 4.299 0.576 *

96094 - 5936 73 11.6 3.725 0.550

98630 - 6060 52 10.0 3.043 0.553

99491 4414 83 Leo 5509 96 8.6 5.230 0.785

101206 - 4649 60 7.6 4576 6.750 0.983

102870 4540 5 Bet Vir 6055 48 6.8 6176 6095 6124 6127 3.407 0.516

107705 4708 17 Vir 6040 56 7.8 4.104 0.498

108954 4767 6037 60 5.5 6060 6068 6068 4.507 0.518

109358 4785 8 Bet CVn 5897 72 6.2 5879 5867 4.637 0.549

110833 - 5075 80 3.9 6.130 0.938

110897 4845 10 CVn 5925 68 12.3 5795 5862 4.765 0.510

112758 - 5203 83 8.4 5116 5137 5.931 0.791

114710 4983 43 Bet Com 5954 71 6.8 6029 5964 5959 5985 4.438 0.546

115383 5011 59 Vir 6012 40 9.3 6021 5989 5967 3.921 0.548

116443 - 4976 83 9.9 6.175 0.850

117043 5070 5610 98 4.7 4.851 0.729

117176 5072 70 Vir 5611 104 4.7 5482 3.683 0.678 *

119802 - 4763 71 6.6 6.881 1.099

122064 5256 4937 84 8.1 6.479 1.038

122120 - 4568 35 11.4 7.148 1.176

124292 - 5535 89 4.0 5.311 0.721

125184 5353 5695 89 5.9 5562 3.898 0.699

126053 5384 5728 79 6.9 5635 5645 5.032 0.600

130322 - 5418 85 5.4 5.668 0.764 *

131977 5568 4683 62 6.8 4605 4609 4551 6.909 1.091

135204 - 5413 91 4.6 5.398 0.742

135599 - 5257 86 5.1 5.976 0.804

137107 5727 2 Eta CrB 6037 60 6.9 4.237 0.507

139323 - 5204 90 7.7 5.909 0.943

139341 - 5242 90 7.9 5.115 0.898

140538 5853 23 Psi Ser 5675 100 3.5 5.045 0.640

141004 5868 27 Lam Ser 5884 81 4.4 5937 5897 4.072 0.558

143761 5968 15 Rho CrB 5865 81 11.1 5782 5726 4.209 0.560 *

144287 - 5414 93 5.7 5.450 0.739

144579 - 5294 89 10.3 5309 5275 5.873 0.707

145675 - 5406 98 12.1 5.319 0.864

146233 6060 18 Sco 5799 96 3.8 4.770 0.614

149661 6171 12 Oph 5294 90 3.2 5.817 0.817

151541 - 5368 88 6.4 5.630 0.757

152391 - 5495 82 4.5 5.512 0.732

154345 - 5503 87 5.6 5.494 0.708

154931 - 5910 82 6.7 3.558 0.578

157214 6458 72 Her 5784 85 9.5 5676 4.588 0.572

157881 - 4035 9 4.5 4011 8.118 1.371

158614 6516 5641 98 3.6 4.910 0.678

158633 6518 5290 83 10.7 5.896 0.737

159062 - 5414 96 7.9 5.485 0.706

159222 6538 5834 93 4.0 5770 5708 5852 4.653 0.617

159909 - 5749 93 5.6 4.459 0.657

160346 - 4983 84 3.9 6.382 0.950

161098 - 5617 90 7.3 5.294 0.632

164922 - 5392 96 6.0 5.293 0.789

165173 - 5505 95 4.7 5.388 0.732

165401 - 5877 85 8.5 5758 4.880 0.557

165476 - 5845 90 5.9 4.406 0.580

166620 6806 5035 75 5.7 4947 4995 4930 6.165 0.871

168009 6847 5826 93 4.0 5781 5833 5826 4.528 0.596

170512 - 6078 43 9.4 3.965 0.542

171067 - 5674 81 6.5 5.191 0.660

173701 - 5423 104 9.7 5.343 0.847

176841 - 5841 92 6.2 4.487 0.637

182488 7368 5435 82 4.4 5.413 0.788

183341 - 5911 85 3.9 4.201 0.575

184385 - 5552 87 4.1 5.354 0.721

184768 - 5713 94 3.9 4.593 0.645

185144 7462 61 Sig Dra 5271 79 6.3 5227 5.871 0.765

186104 - 5753 95 5.8 4.621 0.631

186379 - 5941 67 9.8 3.586 0.512

186408 7503 16 Cyg A 5803 83 3.1 5763 5783 4.258 0.614

186427 7504 16 Cyg B 5752 77 3.5 5767 5752 4.512 0.622 *

187123 - 5824 86 5.0 4.433 0.619 *

187897 - 5887 95 5.0 4.521 0.585

189087 - 5341 83 4.0 5.873 0.782

189340 7637 5816 90 8.4 3.920 0.532

190067 - 5387 100 10.3 5.731 0.707

195005 - 6075 51 6.7 4.302 0.498

197076 7914 5821 75 5.6 5761 5774 5815 4.829 0.589

199960 8041 11 Aqr 5878 78 5.9 5813 4.089 0.590

201091 8085 61 Cyg 4264 17 12.4 4323 7.506 1.158

201092 8086 61 Cyg 3808 5 26.4 3865 8.228 1.308

202108 - 5712 82 7.2 5635 5.186 0.610

203235 - 6071 52 8.4 3.606 0.468

204521 - 5809 74 13.6 5.245 0.545

205702 - 6020 50 4.7 3.839 0.513

206374 - 5622 89 5.4 5.304 0.674

210667 - 5461 81 5.6 5.470 0.800

211472 - 5319 91 5.3 5.835 0.802

215065 - 5726 95 9.7 5.131 0.594

215704 - 5418 95 4.9 5.500 0.795

217014 8729 51 Peg 5778 92 5.4 5755 4.529 0.615 *

219134 8832 4900 63 7.9 4785 6.494 1.009

219396 - 5733 91 5.3 3.918 0.654

220182 - 5372 94 4.7 5.661 0.788

221354 - 5295 95 5.5 5.610 0.830

+32 1561 - 4950 82 6.2 6.493 0.919

+46 1635 - 4273 12 4.2 7.895 1.367

Sun - 5777 889 0.9 4.790 0.65

Contents

High precision effective temperatures for 181 F-K dwarfs from line-depth ratios^,

1 Introduction

2 Observations and temperature calibrations

3 Results and discussion

4 Conclusion

References

5 Online Material

Contents

High precision effective temperatures for 181 F-K dwarfs from line-depth ratios,

1 Introduction

2 Observations and temperature calibrations

3 Results and discussion

4 Conclusion

References

5 Online Material

High precision effective temperatures for 181 F-K dwarfs from line-depth ratios^,